A&A 556, A99 (2013) Astronomy DOI: 10.1051/0004-6361/201220495 & c ESO 2013 Astrophysics

Search for cold and hot gas in the ram pressure stripped Virgo dwarf galaxy IC 3418,

P. Jáchym1,2,J.D.P.Kenney2,A.Ružickaˇ 1,M.Sun3, F. Combes4, and J. Palouš1

1 Astronomical Institute, Academy of Sciences of the Czech Republic, Bocníˇ II 1401, 141 00 Prague, Czech Republic e-mail: [email protected] 2 Department of Astronomy, Yale University, 260 Whitney Ave., New Haven, CT 06511, USA e-mail: [email protected] 3 Eureka Scientific, 2452 Delmer Street Suite 100, Oakland, CA 94602, USA e-mail: [email protected] 4 Observatoire de Paris, LERMA, 61 Av. de l’Observatoire, 75014 Paris, France e-mail: [email protected] Received 4 October 2012 / Accepted 30 April 2013

ABSTRACT

We present IRAM 30 m sensitive upper limits on CO emission in the ram pressure stripped dwarf Virgo galaxy IC 3418 and in a few positions covering H ii regions in its prominent 17 kpc UV/Hα gas-stripped tail. In the central few arcseconds of the galaxy, we 6 report a possible marginal detection of about 1 × 10 M of molecular gas (assuming a Galactic CO-to-H2 conversion factor) that could correspond to a surviving nuclear gas reservoir. We estimate that there is less molecular gas in the main body of IC 3418, by at least a factor of 20, than would be expected from the pre-quenching UV-based star formation rate assuming the typical gas depletion timescale of 2 Gyr. Given the lack of star formation in the main body, we think the H2-deficiency is real, although some of it may also arise from a higher CO-to-H2 factor typical in low-metallicity, low-mass galaxies. The presence of H ii regions in the tail of IC 3418 6 suggests that there must be some dense gas; however, only upper limits of <1 × 10 M were found in the three observed points in 7 the outer tail. This yields an upper limit on the molecular gas content of the whole tail <1 × 10 M, which is an amount similar to the estimates from the observed star formation rate over the tail. We also present strong upper limits on the X-ray emission of the stripped gas in IC 3418 from a new Chandra observation. The measured X-ray luminosity of the IC 3418 tail is about 280 times lower than that of ESO 137-001, a spiral galaxy in a more distant cluster with a prominent ram pressure stripped tail. Non-detection of any diffuse X-ray emission in the IC 3418 tail may be due to a low gas content in the tail associated with its advanced evolutionary state and/or due to a rather low thermal pressure of the surrounding intra-cluster medium. Key words. galaxies: individual: VCC 1217/IC 3418 – galaxies: clusters: individual: Virgo – galaxies: evolution – galaxies: ISM – ISM: kinematics and dynamics – methods: observational

1. Introduction from gas-rich dwarf galaxies which lose their interstellar gas and subsequently quench their star formation because of ram The evolution of galaxies in clusters is driven by interactions be- pressure stripping (Boselli et al. 2008). Objections to this sce- tween the interstellar medium (ISM) of the galaxies and the gas nario based on differences in kinematics, metallicity, and nuclei in the intra-cluster medium (ICM; Gunn & Gott 1972; Poggianti in the dI and dE populations have been largely convincingly ad- et al. 1999; van Gorkom 2004; Koopmann & Kenney 2004). dressed by Kormendy & Bender (2012). But until recently there Because of its proximity and dynamical youth (Schindler 1999; has not been any clear case of a galaxy actually undergoing such Randall et al. 2008), the Virgo Cluster is an excellent laboratory a transformation. for studying the processes that transform galaxies in dense en- vironments. By far the most numerous galaxies in Virgo are the The galaxy IC 3418 (VCC 1217) is a peculiar dwarf galaxy dwarfs. The relationship between the blue star-forming dwarf ir- located very close to the cluster core, with a projected distance of only 1◦ (∼290 kpc) from M87 (see Fig. 1), and a line-of-sight regulars (dIs) and the more common dwarf ellipticals (dEs) has −1 been a long-standing puzzle. Dwarf ellipticals are centrally con- velocity with respect to the cluster mean of about –1000 km s . centrated in the cluster, whereas blue star-forming dwarf irreg- It is very likely on its first orbit through the Virgo center, and probably close to its nearest approach to M87. No H i was de- ulars are preferentially located in the outskirts, although their ∼ × 6 masses and structural parameters are similar (Binggeli et al. tected anywhere in the galaxy down to a 3σ limit of 8 10 M 1987). It has therefore been suggested that dEs might transform per beam (Chung et al. 2009); in the main body of the galaxy there is also no Hα emission (Hester et al. 2010). The stel-  lar metallicity of the galaxy was measured from optical spec- Based on observations carried out with the IRAM 30 m Telescope ∼ ± and with the Chandra X-ray Observatory. IRAM is supported by troscopy to be 0.5 0.2 Z (Kenney et al., in prep.). A new / stellar population analysis indicates that star formation in the INSU CNRS (France), MPG (Germany), and IGN (Spain).   galaxy stopped about 250 ± 50 Myr ago in the central 25 ,and CO spectra as FITS files are available at the CDS via anonymous  ftp to cdsarc.u-strasbg.fr (130.79.128.5)orvia that it took ∼<70 Myr for star formation to cease from ∼40 to http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/556/A99 the center (Kenney et al., in prep.). This is consistent with ram Article published by EDP Sciences A99, page 1 of 15 A&A 556, A99 (2013)

Table 1. Parameters of IC 3418 (VCC 1217).

RA, Dec (J2000) 12:29:43.8, +11:24:09 Type IBm a −1 υhelio 170 ± 10 km s ◦ h DM87 projected 1.02 (≈290 kpc) Major/minor diameter b 1.5/1 (≈7.2 kpc/4.8 kpc)h PA, inclinationc ∼45◦, ∼50◦ Apparent/absolute B mag.d 14.85/–16.2 Apparent/absolute K mag.e 12.55/–18.5 FUV, NUV, FUV-NUV f 17.55, 16.69, 0.85 ± 0.10 g 6 H i mass <8 × 10 M (3σ) H i-deficiencyg >2.16 IC3418 Tail projected length 3.5 (≈17 kpc)h

References. (a) Kenney et al. (in prep.); 38 km s−1 in GoldMine (Gavazzi et al. 2003); (b) NED ; (c) HyperLEDA (Paturel et al. 2003); (d) Internal and Galactic extinction corrected (Gavazzi et al. 2006). 8 Corresponding B-band luminosity is 4.8 × 10 L,usingMB, = 5.48 (e.g., Mo et al. 2010); (e) GoldMine. Foreground Galactic extinction in the K-band is 0.012 mag (NED). Corresponding K-band luminosity is 8 ( f ) 5.2 × 10 L,usingMK, = 3.28 (e.g., Mo et al. 2010); Gil de Paz et al. (2007); (g) Chung et al. (2009); (h) Assuming Virgo cluster distance of 16.5 Mpc (Mei et al. 2007). Fig. 2. Location of IC 3418 in the (FUV − NUV)− MB color–magnitude diagram of nearby galaxies from the sample of Gil de Paz et al. (2007). FUV − NUV = 0.85±0.10 (Gil de Paz et al. 2007)andmB corrected for internal and Galactic extinction is 14.85 mag (NED), i.e., MB = −16.2. Red dots are elliptical/lenticular galaxies, dark green triangles are early- type spirals, light green triangles are late-type spirals, blue asterisks are irregular and compact galaxies, and black diamonds are galaxies lacking morphological classification. Figure adapted from Gil de Paz et al. (2007). M86 fireballs (Yoshida et al. 2008), elongated streams of young stars with H ii regions and bright UV knots at their heads (Kenney M87 et al., in prep.). These are interpreted as dense, star-forming gas clouds which are accelerated by ram pressure, leaving behind trails of newly-formed stars that are not affected by ram pressure and therefore decouple from the gas. While good examples of ram pressure stripping are known among the Virgo spiral galaxies (e.g., Kenney et al. 2004; Chung IC3418 et al. 2007, 2009; Abramson et al. 2011), none of these galaxies is completely gas stripped and none has such a luminous UV tail, or linear parallel stellar streams, or fireballs as IC 3418. There are already known examples of one-sided tails of young stars 0.5 deg extending from more massive galaxies in more distant, richer clusters (Cortese et al. 2007; Yoshida et al. 2008; Smith et al. Fig. 1. Location of IC 3418 in the Virgo cluster. The background image 2010; Sun et al. 2010; Woudt et al. 2008), where the ram pressure shows the distribution of the 0.5–2.0 keV ICM as observed in X-rays can be 1 to 2 orders of magnitude stronger than in Virgo, but with ROSAT (Credit: S. L. Snowden, http://heasarc.gsfc.nasa. IC 3418 is by far the closest such galaxy known. gov). Silhouette of a GALEX UV image of IC 3418 is displayed at the In Fig. 2, IC 3418 is placed into a (FUV − NUV) vs. position of the galaxy. MB color–magnitude diagram of a sample of nearby galaxies studied by Gil de Paz et al. (2007). It shows that IC 3418 stands outside both the red and blue sequences of galaxies, respectively, pressure suppression of star formation from the outside inwards, and that it occurs right at the boundary between early- and late- acting typically on a timescale of ∼100 Myr. The main stellar type galaxies. This suggests that IC 3418 is being transformed body of the galaxy is symmetric and undisturbed indicating that from a dwarf irregular type to an early-type galaxy (see Kenney a tidal interaction is excluded (Kenney et al., in prep.). et al., in prep.) which is consistent with the recent cease of its A ram pressure stripping scenario for IC 3418 is also sug- star formation due to ram pressure stripping. gested by optical and GALEX UV observations (Chung et al. Numerical modeling has shown that ram pressure stripping is 2009; Hester et al. 2010; Fumagalli et al. 2011) that revealed averyefficient process for removing the gaseous content of clus- a remarkable one-sided tail extending 17 kpc from the galaxy’s ter galaxies (e.g., Vollmer et al. 2001; Roediger & Hensler 2005; SE side. It comprises about nine bright knots and linear, paral- Jáchym et al. 2007, 2009). Simulations also predict that material lel filaments of young UV-bright stars. Only in the outer half of in the gas-stripped tails gets compressed by ram pressure and the tail Hα + [N ii] emission is found, mostly associated with radiative cooling can form new molecular clouds in situ, which

A99, page 2 of 15 P. Jáchym et al.: Search for cold and hot gas in IC 3418

Table 2. List of observed positions.

(1−0) (2−1) RA Dec TON Tsys Tsys B2 (J2000) (J2000) (hr) (K) (K) B1 12:29:43.8 +11:24:09 5.6 240–300 400–500 B2 12:29:44.07 +11:24:22.38 2.8 200–300 230–400 K4 12:29:50.75 +11:22:36.35 2.5 250–350 400–700 B1 K6 12:29:52.29 +11:22:04.86 2.7 250–330 500–600 K8 12:29:53.92 +11:22:01.85 1.7 210–280 240–310

∗ Notes. System temperatures are given in TA scale.

K4 stripped gas, and origin of star formation in the tail. The conclu- sion follows in Sect. 9. K8

K6 1 arcmin 2. Observations 2.1. Millimeter integration: IRAM 30 m Using the 30 m telescope operated by the Institut de Radio Fig. 3. Observed positions on a GALEX image (FUV in blue, NUV Astronomie Millimétrique (IRAM) at Pico Veleta, Spain, in in yellow) of IC 3418 with 12CO(1–0) beams (FWHM = 21)ofthe August 2011 we carried out observations towards the main 30 m IRAM telescope displayed – two positions (B1, B2) in the main body of IC 3418 and the brightest Hα/UV regions of its tail body of the galaxy and three (K4, K6, K8) in the outer part of the tail. (see scheme in Fig. 3). We used the EMIR receiver in E090 The CO(2–1) beams have the same centers but half the diameter as the and E230 bands to observe simultaneously at the frequencies circles shown – a CO(2–1) beam is shown in B1 with a dashed circle. 12 = 12 The background image is adopted from Hester et al. (2010). of the CO(1–0) (νrest 115.271 GHz) and the CO(2–1) (νrest = 230.538 GHz) lines. At these frequencies the telescope half-power beamwidths are 21 and 11, respectively, which might then form a population of new stars (Kapferer et al. 2009; corresponds to a resolution of about 1.7 kpc and 0.9 kpc, respec- Tonnesen & Bryan 2009). Dwarf galaxies have weaker potential tively, at the adopted distance of the Virgo cluster of 16.5 Mpc wells and their ambient ISM may be correspondingly less dense (Mei et al. 2007). As backends, we used the 4 MHz filter- −1 (Bolatto et al. 2008). This, plus a strong ram pressure near the banks at 115 GHz with 10.4 km s velocity resolution, and the cluster center, could significantly enhance stripping of IC 3418 WILMA autocorrelator with a spectral resolution of 2 MHz at −1 and thus account for its peculiar morphology. The presence of both 115 GHz and 230 GHz (i.e., 5.2 km s velocity resolution −1 star formation in the tail, strongly suggesting the presence of at 115 GHz and 2.6 km s at 230 GHz). The new FTS spec- molecular gas, makes IC 3418 an attractive place to search for trometer with 0.192 kHz spectral resolution was also connected CO emission in the environmentally affected galaxy. However, to both lines as a back-up. In this work, WILMA and 4MHz data dwarf irregulars are usually fainter in CO than large spirals when are analyzed. normalized by (stellar) mass and those with metallicities less All observations were done in the Wobbler switching mode than ∼8.0 (in oxygen abundance scale) are challenging to detect that is known to provide much flatter baselines, with ripples at (e.g., Taylor et al. 1998; Komugi et al. 2011). a much lower amplitude than with the position switching mode. The secondary reflector was switching symmetrically at 0.7 kHz In this paper we want to address two main questions: (1) Is  the main body of IC 3418 totally stripped or did some gas survive frequency with 120 amplitude in azimuth. The actual distance between the source and a reference position is double the ampli- the stripping? We present the results of our IRAM 30 m search  for CO emission since molecular gas is more likely to survive tude value, i.e., 240 , which is large enough to just avoid the tail than H i. (2) How does the ISM behave in ram pressure stripped, with the standard azimuthal direction of switching. star-forming tails? We are interested in searching for molecular The main beam efficiencies Beff of the 30 m antenna are esti- and for hot X-ray emitting gas in the tail in order to help char- mated to 0.79 at 115 GHz and 0.57 at 230 GHz, and the forward acterize the ISM in such tails. We present the results of our new efficiencies Feff to 0.94 and 0.92, respectively. The corrected an- sensitive Chandra observation of the galaxy. tenna temperatures provided by the telescope were converted to = ∗ The structure of the paper is as follows: after a brief descrip- the main beam brightness temperature by Tmb TA Feff/Beff.To tion of our observations (Sect. 2), we present and analyze our check the system setup we observed IRC+10216 and M99 and IRAM 30 m results (Sect. 3), and estimate upper limits on the we obtained lines matching the catalogued lines. current molecular gas content of the whole IC 3418 (Sect. 4). We Weather conditions were typical for the summer afternoon then estimate the original molecular gas content of the galaxy period at Pico Veleta with the PWV exceeding 10 mm in the first from the star formation rate (SFR) in the main body and compare half of the run. Then, after a passage of a weather front, the con- the molecular content of IC 3418 with other low-mass galaxies ditions improved substantially and PWV decreased to 5−7 mm. (Sect. 5). The H2-deficiency of the galaxy is also discussed. In The atmosphere was, however, rather unstable and pointing and Sect. 6 we analyze the Chandra results. Then, by means of nu- focusing were difficult. Pointing was checked about every 1.5 h merical calculations we study the effects of ram pressure on ISM mostly on Saturn which was close to our sources, usually with with different column density (Sect. 7). In Sect. 8 we discuss small corrections from 2 to 5. The system temperatures are the structure of the gas-stripped tail, as well as the fate of the given in Table 2.

A99, page 3 of 15 A&A 556, A99 (2013)

Table 3. Properties of the observed positions in IC 3418.

CO(1−0) CO(2−1) Source rmsW(1−0)/W(2−1)/4MHz FWHM ICO(1−0) ICO(2−1) LCO(1−0) LCO(2−1) Mmol Mmol −1 −1 −1 5 −1 2 5 −1 2 6 6 (mK) (km s )(Kkms)(Kkms)(10Kkms pc )(10Kkms pc )(10M)(10M) (1) (2) (3) (4) (5) (6) (7) (8) (9) B1 1.1 / 1.6 / 1.0 54 ± 12 0.06 ± 0.03 0.24 ± 0.05 2.12.10.91.2 B2 1.4 / 2.0 / 1.3 ∼30 <0.1 <0.1 <2.7 <0.9 <1.2 <0.5 K4 2.0 / 3.6 / 1.8 ∼30 <0.1 <0.2 <3.8 <1.7 <1.7 <0.9 K6 2.0 / 4.0 / 1.6 ∼30 <0.1 <0.2 <3.7 <1.9 <1.6 <1.0 K8 1.7 / 2.3 / 1.7 ∼30 <0.1 <0.1 <3.2 <1.1 <1.4 <0.6

Notes. The central velocity of the line in the B1 position is 103 km s−1. Column (1) source name. Column (2) noise levels in WILMA CO(1–0), WILMA CO(2–1), and 4MHZ spectra with 10.4 km s−1 channels. From fitting first-order baselines in a 2000 km s−1 wide region of the spectra. Column (3) linewidths used for both CO(1–0) and CO(2–1) lines. Measured in B1 from the CO(2–1) detection, estimated for other positions (see the text). Column (4) detected (B1) or 3σ upper limits on WILMA CO(1–0) integrated intensity. Column (5) detected (B1) or 3σ upper limits on CO(2–1) integrated intensity. Column (6) detected (B1) or 3σ upper limits on CO(1–0) luminosity. Column (7) detected (B1) or 3σ upper limits on CO(2–1) luminosity. Column (8) detected (B1) or 3σ upper limits on molecular gas mass, including a factor of 1.36 to account for the effect of He, corresponding to WILMA CO(1–0) luminosity. Column (9) detected (B1) or 3σ upper limits on molecular gas mass, including a factor of 1.36 to account for the effect of He, corresponding to WILMA CO(2–1) luminosity.

The sky coordinates of the observed positions shown in suggested from the existence of H ii regions (Fig. 3, see coordi- Fig. 3 are given in Table 2, together with respective on-source nates in Table 2). The B1 pointing is centered at the galaxy’s op- times. These were typically 2.5 h per position per polarization, tical center (as given by NED), and the B2 point is shifted by less divided into 6 min scans. For the B1 position it was roughly than one CO(1–0) beamwidth in the NE direction, towards the double. For K8, the on-source time was shorter by a factor of brightest region of the galaxy where significant substructure oc- ∼1.5 than for the K4 and K6, but observing conditions had curs, indicating that star formation occurred there recently. The improved resulting in a better sensitivity. The data were re- tail positions K4, K6, and K81 are associated with peaks in UV duced with CLASS (Continuum and Line Analysis Single-dish and Hα emission. Software) developed by IRAM in the standard manner: the spec- tra were checked for errors (spikes) and bad channels exceeding the 5σ level were flagged. The quality of WILMA and 4 MHz 3.1. Main body data was very good. Linear baselines were fitted. Each spectrum − First we searched for CO emission in the central position B1. was re-binned to a velocity resolution of 10.4 km s 1 and both This was our deepest integration with the total on-source time polarizations were averaged in order to increase the signal-to- of about 5.6 h (per polarization). The left panels in Fig. 4 show noise ratio (S/N). All scans were then summed up to produce WILMA CO(2–1) and WILMA CO(1–0) spectra smoothed to the final spectra. The root mean square (rms) noise levels of 10.4 km s−1 resolution. Very low rms values of about 1.7 mK and about 1−2 mK at CO(1–0) frequency were achieved in the final 1.1 mK, respectively, were achieved. Although no strong detec- summed spectra (Table 3). tion appears, a feature that is six channels wide (∼60 km s−1) occurs at both frequencies at the central velocity of about −1 2.2. X-ray imaging: Chandra 100 km s . It is more prominent in the CO(2–1) spectrum. We inspected the surrounding parts of the spectra and found no The observation of IC 3418 was performed with the Advanced other similar case. The S/Ns of the line feature in the WILMA CCD Imaging Spectrometer (ACIS) on November 12, 2012 (ob- CO(1–0) and WILMA CO(2–1) spectra are 2.1 and 5.6, respec- = −1/2 −1 N sID: 13811). Standard Chandra data analysis was performed tively, where we have calculated S/N N σrms i=1 Ii, which includes the corrections for the slow gain change, charge where N is the number of channels covered by the line, σrms transfer inefficiency, and the ACIS low-energy quantum effi- is the rms intensity of the spectrum per channel, and Ii is the ciency degradation from the contamination on the ACIS opti- brightness temperature in the ith channel. In order to claim a de- cal blocking filter. No flares of particle background were present tection, we require the integrated intensity ICO = Tmb dυ to in the observation. The effective exposure time is 33.8 ks for be at least three times greater than the noise over the spectral the S3 chip where IC 3418 is positioned. CIAO4.4 was used channels covered by the line2. for the data analysis. The calibration files used correspond to There is no statistically significant line detection in the Chandra calibration database 4.5.3 from the Chandra X-ray CO(1–0) WILMA and 4MHz spectra of the central position, Center. The solar photospheric abundance table by Anders & but there are 2.0–2.7-σ features with the same velocity range Grevesse (1989) were used in the spectral fits. We adopted as the CO(2–1) line. The integrated CO(1–0) intensity measured an absorption column density of 2.2 × 1020 cm−2 from the from the WILMA spectrum is ∼0.06 K km s−1 (see Table 3). The Leiden/Argentine/Bonn H i survey (Kalberla et al. 2005). CO(1–0) and CO(2–1) features are consistent with a compact CO source which suffers much more beam dilution in the 21 CO(1–0) beam than the 11 CO(2–1) beam. With four times 3. Results – IRAM 1 The naming of the tail positions corresponds to Hester et al. (2010). In a set of deep integrations we observed two positions in the Fumagalli et al. (2011)useadifferent notation (K-6, K-4, K-3). 2 main body of the galaxy where dense molecular gas might have The noise over N spectral channels covered by the line σI = ff 1/2 1/2 survived the e ects of the cluster environment, and three po- Δυres σrmsN = σrms(Δυres ΔυCO ) ,whereΔυres is the velocity res- sitions in the outer tail where the presence of dense gas is olution and ΔυCO is the mean FWHM linewidth.

A99, page 4 of 15 P. Jáchym et al.: Search for cold and hot gas in IC 3418

Fig. 4. Marginal detection in the central main body position B1 (left), spectra in the off-center position B2 (middle), and combined K4+K6+K8 tail spectra (right): WILMA CO(2–1) (top) and WILMA CO(1–0) (bottom) spectra smoothed to 10.4 km s−1 resolution. Dashed lines show ±1σ noise −1 levels; y-axis scaling is set to ±4σ. The velocity scale is LSR; for the sky position of IC 3418, υhelio = υLSR − 3.2kms . the beam area, a point-like source will experience four times the The size of the emitting region in the B1 pointing is certainly beam dilution in the CO(1–0) beam. The ratio of peak temper- comparable to or smaller than the CO(2–1) beam, i.e., ∼<900 pc. atures of ∼3 and integrated intensities of ∼4 are consistent with This is most likely gas concentrated in the nuclear region. the factor of 4 expected from maximum beam dilution multiplied A lower limit on the column density of the central CO-emitting −2 by the typical ICO(2−1)/ICO(1−0) line ratio of 0.8–1 in molecular blob is ∝ Mmol/ΩB ∼ 1 M pc , assuming a standard Galactic clouds (e.g., Leroy et al. 2008). conversion X-factor, if gas uniformly fills the beam. It is, how- We thus consider the CO(2–1) emission in the center of ever, much more likely to be concentrated over a smaller area. IC 3418 as marginally detected. In Fig. 4, the CO(2–1) spec- The actual gas surface density could thus be similar to that of −2 tral line is fitted with a Gaussian profile with the central ve- a typical giant molecular cloud (GMC), or ∼100 M pc .Also − locity of 103 km s−1,FWHMof54kms−1, peak brightness the B1 luminosity of 2.1 × 105 Kkms 1 pc2 is comparable to the temperature of 4.3 mK, and the integrated intensity ICO(2−1) = luminosity of a single giant Galactic molecular cloud. However, 0.24 ± 0.05 K km s−1. The central velocity of the line is some- the linewidth is much larger than that from a single GMC with 6 ∼ −1 what smaller than that deduced by Kenney et al. (in prep.) from Mgas 10 M (typically 10 km s ) which strongly suggests the stellar Keck spectra (see Table 1). We discuss the veloc- that emission is not from a single GMC.  ity difference in Sect. 8.1. Since there is no H i emission in In the off-center main body position B2 (about 15 ∼ 1.2 kpc IC 3418, we estimate the maximum rotation velocity from the from B1) covering the NE young stellar complex (which is the Tully-Fisher relation (by comparing the H-band luminosity of brightest optical part of the galaxy), no CO emission was de- IC 3418 to other Virgo galaxies with known rotation velocities) tected (see Fig. 4, middle panels). This suggests that molecular to be about 55–60 km s−1. Correcting for an inclination of ∼50◦ gas has possibly survived the effects of the cluster environment (Chung et al. 2009), the expected linewidth is about 50 km s−1, only in the very center of the galaxy. In the B2 position, we as- − similar to the observed linewidth. sume an upper limit on the line width to be about 30 km s 1, At the adopted Virgo distance of 16.5 Mpc, the main beam i.e., smaller than at B1, as it does not cover the rotation cen- 2 projected area at CO(2–1) frequency is ΩB 137 arcsec = ter. Rotation curve studies of Local Volume dwarfs (e.g., Kirby 0.9 kpc2, including a Gaussian beamshape correction factor of et al. 2012) show that at r = 1 kpc, the rotation velocity is ∼ −1 1/ ln 2. The B1 integrated intensity then corresponds to a lu- 0.5υmax 30 km s . Thus, we derive the upper limits on the 5 −1 2 minosity LCO(2−1) = 2.1 × 10 Kkms pc . The correspond- integrated CO intensity, CO luminosity, and molecular gas mass, = given in Table 3. ing molecular gas mass can be calculated using MH2 [M] R21 −1 2 = 5.5 XCO 0.8 LCO [K km s pc ], where XCO NH2 /ICO is the CO(1–0)-to-H2 conversion factor normalized to a standard 3.2. Stripped tail and fireballs Galactic value of 2 × 1020 cm−2 (K km s−1)−1 (e.g., Kennicutt & Evans 2012; Pineda et al. 2010; Feldmann et al. 2012), and R21 We searched for CO emission in the three points K4, K6, and is the CO(2–1) to CO(1–0) ratio. We assume a typical value of K8 covering most of the bright H ii regions in the outer tail R21 = 0.8(e.g., Leroy et al. 2008). The above formula includes (see Fig. 3). The positions K4 and K6 show a head-tail (fireball) a factor of 1.36 to account for the effects of helium. The result- structure with Hα offset from the UV peaks. We did not observe ing amount of molecular gas possibly detected in the B1 main the brightest UV region (K5 in the notation of Hester et al. 2010) 6 body position is 1.2 × 10 M. In Sect. 5.2 we discuss that the and instead went for the more distant knot K8 that has a higher X-factor may be (somewhat) larger than the standard Galactic Hα/UV ratio than K5, suggesting that it may be at an earlier value because of the sub-solar metallicity of IC 3418. evolutionary stage with more gas. The galaxy IC 3418 is moving

A99, page 5 of 15 A&A 556, A99 (2013)

Fig. 5. Tail positions K4, K6, and K8: WILMA CO(2–1) (top) and WILMA CO(1–0) (bottom) spectra smoothed to 10.4 km s−1 resolution. Dashed lines show ±1σ noise levels; y-axis scaling is set to ±4σ. toward us with respect to M87 and the stripped material is thus Table 4. Molecular gas mass estimates corresponding to our IRAM de- extending away from us: we expect it to be accelerated to larger tection and upper limits on CO luminosity in the main body and in the line-of-sight velocities. tail, assuming a Galactic CO-to-H2 conversion factor. In Fig. 5 we inspect the WILMA spectra for the K4, K6, and K8 regions over a large range, from velocities of the main body Source LCO(2−1) Mmol −1 5 −1 2 6 to those several hundred km s higher. Although there are few (10 Kkms pc )(10M) ∼3σ features resembling spectral lines, they never occur in the Center (B1) 2.1 1.2 spectra at both frequencies simultaneously, as in the case of B1. Total disk ≥2.1, <9 ≥1.2, <5 As we will show later in Figs. 9 and 10, only low-column den- Total tail <18 <10 sity ISM parcels could be accelerated by ram pressure to large Body + tail <27 <15 vertical velocities. No clear line emission is thus detected in any of the three tail positions. Since the tail linewidths might be broadened by galactic rotation, we will assume linewidths of − 4.1. Current upper limits – the whole galaxy ∼30 km s 1 to derive upper limits on CO intensity in all three tail positions. Table 3 gives corresponding upper limits on ICO, In Table 4 we give estimates (upper limits) of the CO luminosity LCO,andMmol. The upper limits on the molecular gas mass are and molecular gas mass in the central beam of the main body 6 ∼1 × 10 M per point, assuming a standard Galactic X-factor, (the B1 position), the entire main body, the tail, and the whole but it is possible that in the Hα bright tail regions the CO-to-H2 galaxy (main body plus tail). To estimate the total CO luminos- conversion factor is different from standard value (see discussion ity of the main body of the galaxy from the measured B1 and B2 in Sect. 5.2). (upper limit) luminosities, we estimate the filling factor of the In order to increase the SNR, we also combined data from CO(2–1) and CO(1–0) beams as the ratio of their area to the op- the three positions, assuming for a moment that conditions of tical surface of the galaxy, assuming an exponential distribution the molecular gas are similar and that their radial velocities are of CO emission with a scale length roughly 1.5 times smaller close; however, no significant signal occurred in the stacked than the optical scale length (Young et al. 1995). Fumagalli spectrum (see the right panels in Fig. 4). Our achieved upper et al. (2011) fitted the surface brightness of the IC 3418 disk limits per point on CO luminosity are comparable to the lumi- with an exponential law with a scale length of 19 1.5 kpc. ∼ × 5 −1 2 nosity across most of the SMC, LSMC,CO10 1 10 Kkms pc Accounting for the R25 major axis diameter of ∼7.2 kpc and the as measured by Mizuno et al. (2001). With CO(1–0) point lumi- inclination of ∼50◦ (see Table 1), we estimate that the CO(1–0) − nosities of ∼2 × 105 Kkms 1 pc2, we are sensitive enough to de- main beam located in the B1 position would encompass about tect analogues to the brightest clouds in M33 or the LMC N197 25% of the total disk CO luminosity. This yields an estimated complex. upper limit on the molecular gas content in the whole galaxy 6 of about 5 × 10 M. Leroy et al. (2008) suggested that typical dwarfs in a sample of nearby galaxies have molecular gas dis- < 4. Current molecular gas content tributed mostly in their inner ( 0.25r25) radius, which matches the CO(1–0) beam radius. Moreover, since the galaxy is prob- In this section we will estimate from our IRAM observations up- ably gas-stripped from the outside in, the expected gas distri- per limits on the molecular content in the whole galaxy. We will bution would be more compact than the stellar disk. Thus, it is also calculate an upper limit on the molecular gas mass follow- likely that we are detecting a larger fraction of the total CO flux ing from dust observations of IC 3418. in the central beam.

A99, page 6 of 15 P. Jáchym et al.: Search for cold and hot gas in IC 3418

6 Applying our 3σ upper limit of Mmol ∼< 1 × 10 M mea- estimate that FUV flux in a case where star formation ended sured in each of the three tail regions (K4, K6, and K8) to the 100 Myr or 200 Myr ago is reduced by a factor of 8 or 20, total number of knots occurring in the wake, an upper limit on respectively, compared to a model with continuous star forma- 6 the total molecular mass in the tail would be ∼9×10 M.Asim- tion (Hugh Crowl, priv. comm.). This factor could be somewhat ilar value comes from calculating the filling factor of the three lower if ram pressure induced a burst of star formation before CO(1–0) beams in the whole tail area. Assuming that the molec- quenching it, a phenomenon that was identified in the outer radii ular gas occurs (mostly) in the actual places of star formation, of NGC 4522 (Crowl & Kenney 2006), for example. For IC 3418 and taking into account that all H ii regions are observed only this is suggested from measuring equivalent widths of the ab- in the outer half of the tail, the upper limit could possibly be sorption Balmer lines in the spectrum of IC 3418, which are lowered to a value similar to the one derived above for the main stronger than in post star-burst k+a galaxies (Fumagalli et al. body. Nevertheless, we conclude that the upper limit on the cur- 2011). rent molecular gas content of the whole galaxy (main body plus As a result, we will apply a factor of ∼10 to the SFR of 7 −3 −1 tail) is about 1.5 × 10 M. 6.1 × 10 M yr determined from the FUV luminosity–SFR formula of Kennicutt & Evans (2012)     4.2. Current molecular gas amount from dust emission −1 −1 log SFR M yr = log LFUV erg s − 43.35. (2) Far infrared (FIR) dust measurements allow an alternative ap- proach for estimating the current molecular gas mass that To account for sub-solar metallicities appropriate for dwarf overcomes some disadvantages of traditional CO line tracers. galaxies, we will also apply to the SFR the factor of 1/1.1 It traces the total gas column density (H i+H2). In the Small estimated by Hunter et al. (2010) from the stellar population evolution models of STARBURST99 (Leitherer et al. 1999) Magellanic Cloud, IRAS observations suggested much more H2 than seen from CO (Israel 1997; Leroy et al. 2007a) and, thus, for luminosity at 1500 Å. This then yields an original SFR of −2 −1 ∼ . × M that CO may be underabundant or absent in regions where H2 5 6 10 yr in IC 3418, a value that is consistent with survives because the H2 self-shields while the CO is photo- measurements in other dwarf galaxies where SFRs can span sev- dissociated (e.g., Maloney & Black 1988). eral orders of magnitude (e.g., Hunter et al. 2010; Schruba et al. From ISO observations Tuffsetal.(2002) obtained 3σ upper 2012). Assuming a typical star formation efficiency in dwarf ≈ limits on 60 μm, 100 μm, and 170 μm FIR dust emission of galaxies of τdep 2Gyr(e.g.,Bigiel et al. 2008), this yields 8 IC 3418 of <0.04 Jy, <0.03 Jy, and <0.07 Jy, respectively. The about 1 × 10 M of molecular gas that originally was in the total dust mass can be determined to within a factor of 3 from galaxy. This is about 20 times more than our CO-based upper (e.g., Boselli et al. 2002; Evans et al. 2005) limit estimate of the current molecular gas content of the main   body. 2 a/Tdust Mdust = CS100 μm D e − 1 , (1) Another rough estimate of the original molecular amount can come from the stellar mass of the galaxy and typical fractions where C relates to the grain opacity, S 100 μm is the flux in the of atomic and molecular gas seen in galaxies of the same type. 100 μmbandinJy,D is the distance of the galaxy in Mpc, and However, this estimate is affected by the issue of the X-factor Tdust is the dust temperature. According to Boselli et al. (2002), that we will deal with later. Bothwell et al. (2009) calculated the = −1 −2 = = we use C 1.27 M Jy Mpc , Tdust 20.8K,anda 144 K stellar mass-to-light ratios for different galaxy morphology-color for 100 μm emission. For the S 100 μm < 0.03 Jy upper limit we combinations using the algorithm of Bell & de Jong (2001). For 4 get a limit on the dust mass of ∼<10 M. Taking a typical cold Im and dIrr types they derived M/LB = 0.49. For the B-band gas-to-dust ratio of about 500 (e.g., Lisenfeld & Ferrara 1998), magnitude mB = 14.85 (see Table 1) of IC 3418 this yields a 8 we obtain an upper limit on the molecular gas mass in the galaxy stellar mass M∗ ≈ 3 × 10 M, which is consistent with the es- 6 8 of ∼<5 × 10 M. This is a value similar to our CO luminosity- timate of Fumagalli et al. (2011)(∼3.8 × 10 M). It is known based upper limit estimate (see Table 4). As we will discuss in that ISMs of dwarf galaxies are dominated by large reservoirs the following section, it is likely that molecular gas has been of atomic gas – along the Hubble sequence the H i fraction in- ff e ectively stripped from IC 3418. This would imply that some creases up to typically MH i/M∗ > 0.5oreven>1 in the late types dust also has been removed from the galaxy by ram pressure (e.g., Bell & de Jong 2000). The original H i mass of IC 3418 is 8 stripping. estimated to ∼6×10 M (Gavazzi et al. 2005),whichisinagree- ment with the expected MH i/M∗ > 1. Although dwarf galaxies are H i-rich, molecular gas represents only a small fraction of 5. Original molecular gas content ≈ their total gas mass, typically MH2 /MH i 0.1–0.2 which is sim- We estimate the molecular gas mass Mmol likely to have been ilar to the ratio found in the outer parts of spiral galaxies (e.g., in IC 3418 before it was ram pressure stripped from the SFR in Obreschkow & Rawlings 2009; Israel 1997). This corresponds 8 the main body. The far-ultraviolet (FUV) emission is a tracer to ∼(0.6−1) × 10 M of original H2 in IC 3418, which is con- of star formation in the last couple 100 Myr. The FUV lu- sistent with the above estimate. minosity within the 25 mag arcsec−2 B-band isophote of the main body of IC 3418 (corrected for Galactic extinction) is 41 −1 5.1. Comparison with other galaxies LFUV = 1.4 × 10 erg s (Gil de Paz et al. 2007). The classical SFR-luminosity relation works well for continuous star forma- We will also compare our results with other galaxies, including tion approximation where SFR is constant over the timescale of low-mass, low-metallicity dwarfs. In Fig. 6, IC 3418 is placed the UV emission. This is probably not true in IC 3418, where into a plot of normalized CO luminosity LCO/LK vs. LK ,to- star formation ceased about 200 Myr ago as a result of ram pres- gether with (a) the sample of Schruba et al. (2012) containing sure stripping of star-forming ISM, and so the current FUV lu- 16 nearby low-mass star-forming HERACLES galaxies as well minosity is likely scaled down because the stellar population had as some more massive HERACLES galaxies and Local Group faded. From evolutionary stellar population modeling we can galaxies, (b) a subset of the sample of compact late-type spirals

A99, page 7 of 15 A&A 556, A99 (2013)

from a lack of observations with CO sensitivity levels similar to ours, this suggests that CO content of IC 3418 is low relative to its other properties. The LCO/LK ratio is systematically lower 10 in galaxies with LK ∼< 10 L, probably mostly because of a different CO-to-H2 relation connected to their low metallicity. But even among low-mass galaxies of a given metallicity, the CO-to-H2 relationship may vary due to variations in the amount of photo-dissociating UV radiation, which is correlated with the SFR. The other dwarfs with similar or lower values of LCO/LK , notably NGC 1569, NGC 2537, or NGC 5253, are all star- burst blue compact dwarfs in which strong photo-dissociation of molecular gas by UV radiation from young stars may make their CO content especially low. In IC 3418, since star formation ended about 100–200 Myr ago, the effects of photo-dissociation are expected to be less extreme. Thus, the reason for the low LCO/LK ratio of IC 3418 must be different. Since the galaxy clearly has been strongly ram pressure stripped, the removal of molecular or molecular-forming ISM would be a possible explanation.

5.2. CO as a tracer of H2 It has been shown that in dwarf irregular galaxies with metallici- Fig. 6. Normalized CO(1–0) luminosity as a function of K-band lu- ties lower than about 1/10 of solar, i.e., 12+log(O/H) ≤ 8.0, the minosity. IC 3418 is represented with a red arrow with the head corre- detection of CO emission steeply decreases (e.g., Taylor et al. sponding to the marginally detected (B1) CO luminosity and the upper 1998; Leroy et al. 2007b; Komugi et al. 2011; Schruba et al. tail bar to the upper limit for the total luminosity of the disk. IC 3418 is 2012), although ionized carbon and the ratio of IR to millimeter compared with a set of nearby low-mass star-forming galaxies from the dust emission around star-forming regions indicate that molecu- HERACLES survey (together with more massive HERACLES galax- lar gas is still present. Applying a Galactic CO-to-H2 conversion ies, and other Local Group and nearby galaxies) used by Schruba et al. ratio then results in very low molecular gas masses and about (2012)(filled circles), and with a sample of compact late-type spiral and ffi irregular galaxies of Leroy et al. (2005)(empty circles), plus few star- 10–100 times higher star formation e ciencies (SFEs). It is, burst dwarf galaxies from Taylor et al. (1998)(triangles), and DDO154 however, more likely that the true SFEs of dwarfs are normal, ff (Komugi et al. 2011). Upper limits are plotted with symbols with small and the CO-to-H2 factor is systematically di erent. Thus, assum- arrows. Colors represent metallicity 12 + log (O/H). The color of the ing a constant H2 depletion time of τdep = 1.8Gyr,Schruba et al. IC 3418 arrow does not correspond to its metallicity (see Sect. 5.2). (2012) derived for dwarf galaxies with metallicities 1/2–1/10 Z Small empty circles show marginal detections of Leroy et al. (2005). conversion factors more than one order of magnitude larger than in massive spirals with solar metallicities. Using evolutionary models for Lick indices in Keck spec- and irregulars of Leroy et al. (2005, their detections correspond troscopic data, stellar metallicity that represents an average over instead to lower limits because they observed only the central the star formation history in IC 3418 is estimated to 0.5 ± 0.2 3 few kpc of each galaxy), (c) a few starburst dwarfs of Taylor solar, i.e., 12 + log(O/H) ∼ 8.4 (Kenney et al., in prep.) . et al. (1998) not contained in the Schruba et al. (2012) collec- Another estimate comes from the stellar mass-metallicity rela- tion, and (d) DDO154 observed by Komugi et al. (2011). tion (Tremonti et al. 2004; Kewley & Ellison 2008; Pettini & ∼ × 8 The normalization L /L used in Fig. 6 should remove Pagel 2004): for the IC 3418 stellar mass of 3 10 M, the pre- CO K + ∼ ∼ most of the galactic mass and size correlations and thus reveal dicted stellar metallicity is 12 log(O/H) 8.3, i.e., 0.4solar. an unbiased state of the CO content of the galaxies. All the Unfortunately, there is no data available on the gas metallicity K-band magnitudes of the galaxies were taken from NED, as of IC 3418 that would reveal the actual amount of metals cur- well as their distances (except for the sample of Schruba et al. rently present in the ISM. Nevertheless, in IC 3418 we expect the 2012, for which we used their distances). We preferred to use CO-to-H2 conversion factor to be larger than Galactic because of the decreased dust shielding at a presumably low metallicity in redshift-independent (mostly Tully-Fisher) distances whenever 4 available. The K-band apparent magnitude of IC 3418 is 12.55 the galaxy. From Eq. (5) of Leroy et al. (2013) the correction (see Table 1). The CO(2–1) luminosity detected in the B1 po- to the standard CO-to-H2 conversion factor can be estimated to ∼5. However, this estimate remains highly uncertain. sition was converted to CO(1–0) scale using ICO21 = 0.7 ICO10, to be consistent with Schruba et al. (2012). In Fig. 6 IC 3418 is represented with a red arrow with the 5.3. H2 -deficiency? head corresponding to the detected (B1) CO luminosity and While many Virgo cluster spirals are highly H i-deficient (e.g., the upper tail bar to the upper limit for the total luminosity of Kenney et al. 2004; Chung et al. 2007, 2009; Abramson et al. the disk. When first focusing on the high end of the IC 3418 2011), their molecular gas content is relatively normal or only range, the galaxy occurs in the middle of other galaxies and thus shows no special value of LCO/LK . However, when taking into 3 Weassumeasolarvalueof12+ log(O/H) = 8.69 (Asplund et al. account only the actual detection in the nuclear region of the 2009). 4   disk, IC 3418 shifts to a rather sparsely populated region of the cCO,dark ≈ 0.65 exp [0.4/(D/G) ], where (D/G) is the dust-to-gas plot where only few galaxies have similar or lower values of ratio normalized to the Galactic value of 0.01 (see Sect. 4). A typical −2 the normalized CO luminosities. Although this region may suffer surface density of a GMC of 100 M pc is adopted.

A99, page 8 of 15 P. Jáchym et al.: Search for cold and hot gas in IC 3418 slightly deficient (e.g., Kenney & Young 1986, 1989). This is Table 5. Upper limits on the X-ray emission from the IC 3418 tail. largely a difference between the inner H2-dominated, and outer H i-dominated disks. In galaxies which are clearly ram pressure −1/2 1/2 Model T/ZL0.5−2 Lbol ne fX MX fX stripped, nearly all the gas, including molecular gas, and dust 38 38 −3 −3 7 (keV/solar) (10 erg/s) (10 erg/s) (10 cm )(10M) beyond some gas truncation radius is gone (Vollmer et al. 2012; Cortese et al. 2010). While decoupled molecular clouds have 0.3/0.1 3.5 13 4.3 3.7 been observed beyond the main gas truncation radius in some 0.3/1.0 3.5 7.8 1.6 1.4 0.5/0.1 3.1 7.2 3.1 2.7 galaxies (NGC 4402, NGC 4522, NGC 4438; Crowl et al. 2005; / Vollmer et al. 2008; Vollmer 2009, respectively), the mass in 0.5 1.0 3.0 4.5 1.2 1.1 0.7/0.1 2.9 6.1 2.7 2.4 such features is small and they apparently do not survive long. 0.7/1.0 2.7 4.0 1.1 0.97 Thus, ram pressure stripped spirals in Virgo must be somewhat 5 H2-deficient . Notes. L0.5−2: rest-frame 0.5–2 keV luminosity. Lbol: X-ray bolometric The IC 3418 galaxy is highly H i-deficient, with the defi- luminosity. ne: electron density. fX: X-ray volume filling factor. MX: ciency parameter defH i = log (MH i,orig/MH i,obs) ∼> 2.16 (Chung X-ray mass. et al. 2009; Gavazzi et al. 2005). The above estimated original 8 amount of molecular gas in the galaxy of ∼1×10 M in compar- 6 ison with our CO-based upper limit of 5 × 10 M indicates that IC 3418 could also be poorer in molecular gas by a factor of 20 As shown in Table 5, the derived 0.5–2 keV luminosity limit = ∼ is ∼280 times and ∼100 times lower than the luminosities of (i.e., defH2 log (MH2,orig/MH2,obs) 1.3). The H2-deficiency is generally small in Virgo spirals since ram pressure is not strong ESO 137-001 and ESO 137-002, respectively, the spiral galaxies enough to strip massive galaxies deeply. However, lower mass in a more distant cluster with prominent ram pressure stripped galaxies with shallower potential wells are expected to be more tails with strong X-ray emission (Sun et al. 2010). The X-ray × 7 1/2 completely stripped in Virgo, and are expected to be significantly mass limit is less than 3.7 10 fX M,wherefX is the fill- H2-deficient; IC 3418 may be such a galaxy. While we cannot ing factor of the X-ray emitting gas in the tail. If the X-ray tail presently determine how much of the weakness in CO emission in IC 3418 is cooler than the 0.3 keV assumed in our spectral is due to H2-deficiency vs. a non-standard CO-to-H2 relation, the model, the mass and bolometric luminosity limits would become lack of ongoing star formation in IC 3418 strongly suggests the weaker. galaxy is deficient in molecular gas. In numerical simulations the amount of X-ray emitting gas in the tail was found to depend strongly on the surrounding ICM pressure (Tonnesen et al. 2011). The thermal pressure of the 6. Results – Chandra ICM at the location of IC 3418 in Virgo is PICM = nICMT ∼ −12 −2 −4 −3 2.8 × 10 dyne cm ,wherenICM = 7.1 × 10 cm is the No diffuse X-ray emission is detected from IC 3418 and its tail. local ICM density (see Sect. 7 for details), and kBT = 2.5keV There is also no detection of a shock front ahead of IC 3418. is the local ICM temperature (Shibata et al. 2001). The simula- The Chandra data show three point sources in the tail region, tion run of Tonnesen & Bryan (2010) with a close ICM thermal which will be discussed elsewhere. In this paper, we focus on pressure value of 1.76×10−12 dyne cm−2 suggests that some soft  ×  the limit of the X-ray emitting gas in the tail region (3.5 1 ,or (0.5–2 keV) X-ray emission could be observable in the tail of × 16.8 kpc 4.8 kpc). A 3σ upper limit on the X-ray enhancement IC 3418 for our sensitive Chandra surface brightness limit of over the Virgo ICM background emission at the tail region is de- ∼4 × 10−7 erg s−1 cm−2. This follows from Fig. 2 (right panel) rived from the 0.5–2 keV image. A spectral model of the X-ray in Tonnesen et al. (2011), assuming the same column density of tail is required to convert the count rate limit to the flux limit about 1020 cm−2 for the hot gas in the tail of IC 3418 as is mea- and the mass limit. The X-ray tail is expected to have multiple sured in the simulations. However, the gas mass and gas den- temperature components. However, low statistics of X-ray tails sity in the tail are expected to strongly evolve over time, and so far have hindered a detailed study, so often a single tempera- Tonnesen & Bryan (2010) did not study the evolutionary stage ture component is assumed (see detailed discussions in Sun et al. after the galaxy is fully stripped, like in IC 3418. 2010). The known X-ray tails have T /T ratios of 3–8 (Sun ICM tail Once the main body of the galaxy has been (nearly) com- & Vikhlinin 2005; Sun et al. 2010; Wezgowiec ˙ et al. 2011). If we pletely stripped, the source of new tail gas is depleted. In one simply take this range for the ratio, the expected tail temperature of the simulations of Kapferer et al. (2009),whouseadif- of IC 3418 is 0.3–0.8 keV, for a surrounding ICM temperature ferent numerical hydrodynamics scheme than does Tonnesen of ∼2.5keV(Shibata et al. 2001). In this work, we assume a (smoothed particle hydrodynamics, SPH, vs. adaptive mesh re- single APEC model with temperatures of 0.3–0.7 keV and abun- finement, AMR), a model massive galaxy is completely stripped dances of 0.1 solar or 1 solar (see Table 5). The assumed X-ray by a strong ram pressure. Some 400 Myr after the galaxy was tail is approximated by a cylinder 16.8 kpc long with a diameter stripped, there is still a very long bright X-ray tail that is spa- of 4.8 kpc. The derived upper limits are listed in Table 5. While tially coincident with a star-forming tail. However, IC 3418 is it is often assumed that the stripped ISM is metal rich, a one-T a small galaxy, originally with a much smaller amount of gas. fit always results in very low abundance, which is presumably Thus, after the main body has been nearly completely stripped, caused by the mixing of multi-T gas in the tail (see discussions the gas density of the tail is expected to drop. It is possible that in Sun et al. 2010). the X-ray tail lags farther behind the SF tail which is the case 5 in ESO 137-001, for example, where the brightest X-ray region Fumagalli et al. (2009) have recently reported H2-deficient galaxies, although these are anemic galaxies, with low surface densities of H i is 35 kpc from the galaxy. However, the Chandra field of view covers another ∼2.2 after the end of the SF tail, with no sign of and H2 out to large radii. Whereas truncated galaxies are found almost exclusively in clusters and are clearly ram pressure stripped, the ori- the X-ray enhancement. The X-ray surface brightness is about gin of the anemic galaxies, which are found both inside and outside of the same, so the flux limit would be ∼28% higher if a region of     clusters, is unclear, although they may be starved galaxies. 5.7 × 1 is considered (instead of 3.5 × 1 ).

A99, page 9 of 15 A&A 556, A99 (2013)

− − − − Table 6. Parameters of galaxy and cluster models used for semi-analytic about 1820 cm 3 (km s 1)2 = 1.8 × 10 11 dyne cm 2, assum- calculations of ram pressure stripping. ing a deprojected orbital velocity of 1600 km s−1 for IC 3418. The time evolution of the ram pressure exerted on the disk cor- 8 responds to a profile along an orbit of IC 3418 that is consistent Galaxy: Md = 5 × 10 M ad = 1.5 kpc 9 Mh = 7 × 10 M ah = 3 kpc with its observed plane of the sky position with respect to the Virgo cluster center and line-of-sight velocity. The projected tail = × 14 = Cluster: Mvir 1.4 10 M rs 320 kpc direction was assumed to indicate the current direction of mo- δth = 340 Ωm = 0.3 −22 −3 tion in the plane of the sky. We have studied orbits with peri- ρICM,1(0) = 4.4 × 10 kg m rc,1 = 1.2 kpc −23 −3 to apocenter ratios from about 1:5 to 1:20 that are characteristic ρICM,2(0) = 1.4 × 10 kg m rc,2 = 23.7 kpc β = 0.42 β = 0.52 for radial orbits in galaxy clusters (Ghigna et al. 1998; Boselli & 1 2 Gavazzi 2006; Vollmer 2009). As a fiducial orbit we have chosen an orbit that brings IC 3418 currently almost to pericenter, about 275 kpc from the cluster center, with the plane of the sky veloc- Thus, one reason for the weak X-ray (as well as Hα and H i) −1 ity components υx = υy = 820 km s . The total 3D velocity is emission in the tail of IC 3418 could be that it is in an advanced thus ∼1600 km s−1. The gravitational potential of the Virgo clus- evolutionary stage when the gas density of the tail is expected ter was approximated by the spherically symmetric potential of to drop. In the Virgo cluster, however, almost no X-ray tails i the NFW dark matter halo (Navarro et al. 1996) with parameters are observed in other ram pressure stripped galaxies with H specified in Table 6. α or H tails. The possible reasons for our non-detection thus fur- In Figs. 7 and 8, the vertical distance behind the disk to ther include low pressure of the surrounding ICM (for compar- which parcels with different column densities occurring at dif- ison, the ICM thermal pressure at the position of ESO 137-001 ferent radii (up to 5 kpc; the R major axis radius of IC 3418 is . × −11 −2 25 in the A3627 cluster is about 1 8 10 dyne cm , i.e., about ∼3.6 kpc; see Table 1) can get in our model are shown for two six times higher than in Virgo at the position of IC 3418) or a time moments along the ram pressure event: 100 Myr before the lower tail temperature than we assumed in the spectral model, peak and 100 Myr after the peak. In the pre-peak situation, only which would also mean fast cooling (assuming a not-very-small −2 parcels with ΣISM < 15 M pc can be shifted (not necessarily abundance) and thus a short lifetime. stripped) to at least 20 kpc, i.e., to distances comparable to the projected length of the tail of IC 3418. From the nucleus of the Σ ≈ − −2 7. Ram pressure stripping galaxy, parcels with ISM 5 10 M pc can get to such dis- tances. In the post-peak situation, more material is accelerated to Ram pressure stripping is very likely the principal cause of the larger distances – parcels from outer disk radii with ΣISM up to −2 −2 strong gas deficiency of IC 3418 and of its peculiar morphology about 35 M pc , and from inner radii with ΣISM < 15 M pc , (Hester et al. 2010; Fumagalli et al. 2011). Acceleration due to can get to z ∼ 20 kpc. ram pressure basically depends on the column density of the gas Figures 7 and 8 thus suggest that in our simulation low- parcels and the gravitational potential of the galaxy. The shallow density (H i-like) gas is stripped from the galaxy quite easily and potential of the dwarf makes its interstellar matter more suscep- can get to large distances; in the post-peak situation, all the gas −2 tible to ram pressure and so material with higher column den- with ΣISM ∼< 5 M pc gets to z > 100 kpc. Between the inner sity can be affected. Compared to other ram pressure stripped and outer radii, there is a large difference in the column density galaxies in the Virgo cluster, IC 3418 is also projected closer to of parcels that can be moved to a certain distance. Denser gas the cluster center where the ICM density is higher. Although we parcels are thus difficult to strip substantially. don’t know the true de-projected distance, IC 3418 is probably Figures 9 and 10 depict the vertical velocities relative to affected by a stronger ram pressure than other Virgo galaxies. the galaxy to which ISM elements can be accelerated depend- In this section we use semi-analytic calculations to model ing on their column density, for both the pre-peak and post-peak the effects of a time-varying ram pressure on both atomic and stages. Only low-density stripped parcels can reach high veloc- molecular ISM components in a simple manner by taking into ities in the pre-peak phase and the situation does not change account their column density as a critical parameter. In a static much in the post-peak phase where elements with column den- −2 −1 potential of a two-component model of the galaxy consisting sities <15 M pc can occur at about 400 km s relative to the of the disk and halo (the contribution of a bulge is neglected), galaxy. This excludes the possibility that some weak features we follow dynamics of ISM parcels at different disk radii under noted in the spectra of the K4 and K6 tail positions in Sect. 3.2 the influence of an external time-varying face-on force. Every may correspond to a real CO emission. parcel is assigned a column density. For every parcel we calcu- In Fig. 11 we show the time evolution of the stripping radii late the equation of motion in the vertical direction that takes of ISM parcels with different column densities. The stripping into account local ram pressure and the restoring force from the radius of a disk gas component with some ΣISM corresponds to galaxy. The local ram pressure may be expressed as dυ/dt = a radius outside which it is released from the galaxy’s potential 2 −ρICM |υ − υ0| /ΣISM,whereΣISM is the mean column density of (Etot = Ekin + Epot > 0). Figure 11 provides a timescale for the an ISM parcel and (υ − υ0) is the vertical component of its rela- stripping along our fiducial orbit: H i-like parcels (modeled with −2 tive velocity with respect to the surrounding hot gas. The galaxy a Miyamoto-Nagai profile with Σ0 = 10 M pc that yields −2 is modeled with a Miyamoto-Nagai disk and a Plummer halo ΣR=5 kpc = 0.3 M pc ) start to be released from the outer disk (e.g., Binney & Tremaine 2008) with the parameters specified in by a weak ram pressure at cluster outskirts. It takes about 1 Gyr Table 6. before they are removed completely from R = 3 kpc to the center ◦ −2 At the projected distance of IC 3418 from M87 of about 1 , of the galaxy. Denser gas parcels, with ΣISM up to ∼20 M pc , the density of the ICM is ∼7.1 × 10−4 cm−3, assuming a spher- can also be stripped throughout the disk. This happens on shorter ically symmetric and smooth distribution fitted with a double timescales of 100−200 Myr from R = 3 kpc to 0 kpc as higher β-profile (Matsushita et al. 2002). The parameters used are given ram pressure arising only closer to the pericenter is required. For in Table 6. We get an upper limit on the local ram pressure of yet higher values of the column density the stripping radius starts

A99, page 10 of 15 P. Jáchym et al.: Search for cold and hot gas in IC 3418

Fig. 7. Pre-peak stripping: vertical distances that ISM elements can Fig. 9. Pre-peak stripping: vertical velocities to which ISM elements reach with different column densities from radii 0–5 kpc (with 1 kpc can be accelerated with different column densities from radii 0–5 kpc step; from left to right) in our model 100 Myr before the peak ram (with 1 kpc step; from left to right) in our model 100 Myr before the pressure. peak ram pressure.

Fig. 10. Post-peak stripping. Same as in Fig. 9, but 100 Myr after the Fig. 8. Post-peak stripping. Same as in Fig. 7, but 100 Myr after the peak ram pressure. peak ram pressure.

8. Discussion to increase and subsequently approaches the outer disk radius. −2 8.1. Central bulk of molecular gas This happens at ΣISM ∼ 50 M pc which suggests that parcels with densities at this or a higher level can never be stripped from If our marginal detection of CO emission in the B1 pointing is the galaxy by ram pressure alone. real, it corresponds to gas concentrated in the nuclear region. Our simulation results are consistent with the results of spec- The possible survival of dense gas in the galaxy core (central few tral modeling of observations which suggested that star forma- arcseconds = few hundred pc) could either be because it is the tion in the main body of IC 3418 stopped about 200–300 Myr center of the galaxy, the deepest part of the potential well from ago, on a timescale of less than ∼70 Myr from the outside in whichitismostdifficult to remove gas, or because the surface (Kenney et al., in prep.). Some of the denser gas in Fig. 11 is density of molecular gas is higher than the rest of the ISM, or released only after maximum ram pressure, but this is part of both. While some molecular gas in the nucleus has apparently material that had already been pushed out of the disk plane to survived the stripping, there is no ongoing star formation in this the halo region, and will eventually fall back to the disk (see gas. This might be due to a gas density that is below a threshold Jáchym et al. 2009, for details). Although our calculations de- needed for star formation. The detected gas survived ∼200 Myr pend on the adopted values of the free orbital parameters not after the cessation of star formation (Kenney et al., in prep.) that constrained by observations and on the total mass of the galaxy, is presumably due to removal of most of the ISM. the results are consistent with what is expected from other sim- The survival of molecular gas in the very central region of an ulations and existing observations of galaxies in clusters of otherwise gas-stripped dwarf irregular galaxy could be relevant the size of Virgo (Kenney & Young 1989; Boselli et al. 2002; for the formation of nucleated dwarf elliptical galaxies. Many Fumagalli & Gavazzi 2008): it is not possible to directly strip cluster dEs have distinct nuclear star clusters, and it is not un- the dense (molecular) clouds, so they must have been removed derstood why some dEs have them and others do not. A surviv- from IC 3418 in some other way that we will discuss later. ing central gas reservoir may make it possible to form a nuclear

A99, page 11 of 15 A&A 556, A99 (2013)

6 gas expected to be occurring in the tail to be 6 × 10 M.Our observed (3σ) upper limit on CO amount in the tail (see Table 4) is thus consistent with that predicted for a normal gas depletion timescale and normal X-factor. We would need more sensitive CO observations to show that either of these was non-standard. Table 7 summarizes FUV and Hα luminosities for the main body and the three tail regions that we observed in CO, and com- pares the corresponding SFRs. It shows that SFR estimates from the two tracers are in the knots similar, suggesting that indeed the difference in these quantities for the whole tail is because star formation stopped in parts of the tail. Applying our upper limit CO luminosities in the observed regions, we can estimate from the Hα SFRs lower limit values on the CO-to-H2 conver- sion factor α = τdep × SFR/LCO, assuming a constant depletion time τdep = 2 Gyr. The resulting limits are consistent with the −2 −1 −1 Galactic value αCO ∼ 4.4 M pc (K km s ) . Vollmer et al. ff (2012)andBoissier et al. (2012) determined that the star forma- Fig. 11. Time evolution of stripping radii of ISM parcels with di er- ffi ent column densities: H i-like (radius-dependent Miyamoto-Nagai pro- tion e ciency in the gas stripped tails of Virgo galaxies with re- −2 −2 spect to the available amount of atomic gas was at least 10 times file with Σ0 = 10 M pc which yields ΣR=5 kpc = 0.3 M pc ), and −2 constant ΣISM = 10, 15, 25, and 50 M pc (from left to right) in our lower than within the galaxies. However, Vollmer et al. (2012) model. Ram pressure time profile along our model orbit is shown with found normal values of the SFE with respect to molecular con- the dashed line. tent in the extraplanar regions of the galaxies. star cluster even after the galaxy is mostly gas-stripped. Also, as shown in numerical simulations, high-density clouds may lose 8.3. Where is the gas? angular momentum through interaction with the ICM and drift Although dwarf galaxies are gas-rich systems, in IC 3418 prac- toward the center of the galaxy (Tonnesen & Bryan 2009). tically no gas, neither atomic H i nor molecular H , is observed. The heliocentric central velocity of the marginally detected i 2 −1 No H is detected in the body and the tail, with fairly sensi- central CO line is about 100 km s (see Sect. 3.1). From Keck 6 tive upper limits, MH i < 8 × 10 M (Chung et al. 2009). This spectra, Kenney et al. (in prep.) measured a value of 170 ± ∼ i −1 is 1% of the H mass expected in a typical late type dwarf 10 km s for the central velocity of the stars in the galaxy. The galaxy with the optical size or H-band luminosity of IC 3418. galaxy’s rotation is presumably too small to account for such a We have possibly found only a small amount of molecular gas velocity difference. Moreover, acceleration due to ram pressure 6  MH2 10 M in the central 11 of the main body, and only would push the gas toward a higher radial velocity. On the other 7 upper limits (MH2 < 10 M assuming standard CO-to-H2 con- hand, if dense gas clouds can fall back into the galaxy after they version) outside the central region and in the tail. While our condense in the tail and then become decoupled from ram pres- upper limits are not sensitive enough to clearly demonstrate a sure, they would have blue-shifted radial velocities relative to lack of molecular gas in the tail (see Sect. 8.2), there is very the central stellar disk. This behavior was noted in ram pressure likely a large amount of molecular gas missing in the main body. stripping simulations of Kapferer et al. (2009) and could work in Although this is consistent with the observed absence of any IC 3418. Future more sensitive CO observations could resolve ongoing star formation in the body of the galaxy, the question the origin of the observed mismatch between the velocities of arises of what happened to the molecular gas that once occurred the stellar and molecular components. in IC 3418. Calculations in the previous section have shown that low- 8.2. Molecular gas and star formation in the tail density (atomic) gas can easily be completely stripped by ram pressure from a model of IC 3418, which is consistent with the To assess the current molecular content of the tail of IC 3418, we i inquire how much molecular gas is needed for the amount of star sensitive non-detection of any H emission. Higher-density gas Σ ∼ −2 formation in the tail observed. Using the definition of GALEX with ISM up to 50 M pc then can be accelerated from outer AB magnitudes6, the total FUV luminosity of the tail is about radii only, while the densest (molecular) gas stays bound to the 3.6 × 1025 erg s−1 Hz−1 (Fumagalli et al. 2011). Making the as- galaxy. However, stripping by dynamical pressure of the ICM is sumption of a continuous star formation over the last ∼100 Myr, probably enhanced by other processes such as ablation of dense −3 −1 Eq. (2) predicts a SFR of ∼3.4 × 10 M yr . The total Hα lu- clouds disrupted by hydrodynamic instabilities (e.g., Quilis et al. minosity of the knots in the outer half of the tail of about 2000; Kapferer et al. 2009; Tonnesen & Bryan 2009), or shock- 2.4×1038 erg s−1 (Fumagalli et al. 2011) corresponds to a SFR of heating and dissociation of molecular clouds due to supersonic −3 −1 ∼1.3 × 10 M yr over the last ∼10 Myr (Kennicutt & Evans motion of IC 3418 through the Virgo ICM (e.g., Guillard et al. 2012)7. The difference in SFR estimates is not large and may 2012). Moreover, the structure of the ISM in dwarf irregulars can be mainly because Hα emission is restricted to the outer part of be significantly influenced by stellar feedback and often contains −3 −1 ∼ i the tail. Adopting a SFR in the tail of ∼3 × 10 M yr ,and up to kpc-scale shells and holes seen in H (see, e.g., Zhang assuming a typical star formation timescale (or star formation et al. 2012, and references therein), which could facilitate strip- ffi =Σ Σ ∼ ping of the ISM as was shown in numerical simulations of Quilis e ciency SFE(H2) SFR/ H2 )of 2 Gyr (e.g., Bigiel et al. 2008; Leroy et al. 2013) we can deduce the mass of molecular et al. (2000) with holey disks. Thanks to a shallow potential well and low stellar surface densities, the ISM in dwarfs can also have 6 −2 −1 −1 mAB = −2.5log(fν[erg cm s Hz ]) − 48.6. lower average surface densities than in large spirals. Besides 7 −1 −1 log SFR (M yr ) = log LHα (erg s ) − 41.27. these indirect processes, a very important effect may also come

A99, page 12 of 15 P. Jáchym et al.: Search for cold and hot gas in IC 3418

Table 7. Star formation rates in the observed tail regions and corresponding lower limits on local CO-to-H2 conversion factor.

Source LFUV LHα SFRFUV SFRHα LCO αCO 39 −1 37 −1 −4 −1 −4 −1 5 −1 2 −2 −1 −1 (10 erg s )(10erg s )(10M yr )(10M yr )(10Kkms pc )(M pc (K km s ) ) Disk 139.3 – 62.2 – <11 – K4 7.7 8.0 3.5 4.3 <2.6 >3.3 K6 8.6 5.2 3.8 2.8 <2.3 >2.4 K8 4.6 2.7 2.1 1.5 <2.1 >1.4

Notes. FUV and Hα luminosities come from Fumagalli et al. (2011). FUV- and Hα-based SFRs were calculated using the formula of Kennicutt & Evans (2012). For the X-factor estimates, a 2 Gyr H2 depletion time was assumed.

8 1/2 Table 8. Upper limit (3σ) masses (in M) of gas components issued theory, the mass limit of the Hα-emitting gas is 2.8×10 f M, ff Hα from di erent observations of the tail. where fHα is the filling factor of the Hα gas. Numerical simula- tions of ram pressure stripping indicated that bright Hα emis-

H2 H i Hα X-rays sion is produced at the edges of dense neutral clouds (Tonnesen et al. 2011), so the volume filling factor should be low. From M <1 × 107 <4 × 108 <4 × 108 f 1/2 <4 × 107 f 1/2 tail Hα X numerical simulations of Tonnesen et al. (2011) with the value Ref. this work Ch09 this worka this work of ICM thermal pressure consistent with that for ESO 137-001 (their Fig. 2, left panel), it follows that the typical local den- = (a) Notes. Ch09 Chung et al. (2009); BasedonHα surface limit from sity of the Hα gas in the ESO 137-001 tail is ∼1cm−3.Since Kenney et al. (in prep.). the rms electron density observed in the tail of ESO 137-001 is ∼0.045 cm−3 (Sun et al. 2007), the filling factor, defined as from the rather limited lifetimes of dense clouds (<100 Myr) the ratio of the average density and local density, is estimated − that are comparable to the typical timescales of ram pressure to be about 5 × 10 2. Thus, a reasonable upper limit on the stripping. The molecular gas reservoir could not then be replen- amount of Hα emitting gas in the tail of IC 3418 would be about 7 ished by further condensation and molecularization of H i since 6 × 10 M. Considering the atomic component of the gas in 8 that was removed from the galaxy. All these processes, when the tail, the mass upper limit is ∼4 × 10 M following from 6  combined together, may have resulted in effective stripping of the upper limit of 8 × 10 M per beam (FWHM ∼ 16 , Chung molecular ISM in IC 3418. As a result, IC 3418 may currently et al. 2009), and assuming the same tail dimensions used above be virtually deficient of almost all gas except possibly in its very for the X-ray mass estimate (see Sect. 6). The actual amount of center. the H i in the tail may be much lower than the limit since denser clouds have probably condensed into stars while the diffuse gas had been stripped to large distances (very likely exceeding the 8.4. Gas phases in the tail observed length of the tail), as follows from the calculations in Table 8 summarizes the upper limits on the amount of differ- Sect. 7. Numerical simulations of Tonnesen et al. (2011) suggest that in the rather low thermal pressure environment surrounding ent gas phases observed in the tail. Once stripped, the gas has 4 probably been heated to temperatures between that of H i and IC 3418, the mass fractions of cool (300 K < T < 10 K) and × 5 × 7 the ICM, as cooler gas in the tail mixes with hot ICM gas, or hot (7 10 K < T < 4 10 K) components in the stripped gas more dense gas has been condensing into star-forming clouds. might be both of about 30% of the total stripped mass without The absence of bright X-ray emission (see Sect. 6) is likely much evolution as the tail ages from 100 Myr to 250 Myr. because of a low gas content in the tail associated with its ad- / vanced evolutionary state and or because of rather low thermal 8.5. Origin of star formation in the tail pressure of the ICM at the position of IC 3418 in the Virgo cluster. Compression of stripped gas by the surrounding ICM Galaxies with star forming tails are rare in the Virgo cluster. determines the temperature and density distribution of gas in Extraplanar Hα emission was found in only a few ram pres- the tail. If present, the hot gas component in the tail is thus sure stripped galaxies (NGC 4522, NGC 4402, NGC 4330, probably too diffuse to be observable in X-rays. The actual up- NGC 4438, Kenney & Koopmann 1999; Cortese et al. 2003; 7 per limit mass is close to ∼4 × 10 M since the filling factor Abramson et al. 2011; Boselli et al. 2005, respectively) where it of the X-ray emitting gas is not expected to be much smaller typically occurs only fairly close to the disk. Those are galaxies than 1. Associated with the H ii regions, there is clearly ion- with typical (stellar) mass at least a factor of 10 higher than that ized gas in the tail, but it is difficult to reliably estimate its mass of IC 3418. In IC 3418 the situation is different – a large num- without knowing the gas density; however the ionized gas mass ber of young stars occur at quite large distances (up to 17 kpc in H ii regions is generally small. While no diffuse Hα emis- in projection) from the disk distributed in a straight, narrow tail. sion has been detected, there could be significant mass asso- Chung et al. (2009) suggested that recent star formation in the ciated with undetected diffuse 104 K gas. We derive the upper tail occurred from compression of the stripped H i gas. Similarly, mass limit on Hα emitting diffuse gas from the 2σ limit of the Hester et al. (2010) speculated that star formation in the wake Hα surface brightness of 1.36 × 10−17 erg s−1 cm−2 arcsec−2 ob- was driven by turbulence that enhances density contrasts and served with WIYN 3.5 m telescope (Kenney et al., in prep.). thus aids gas cooling and condensation. A similar example of This is ∼2.3 times more sensitive than the Hα image presented star formation in situ in the stripped gas is known in the Virgo by Fumagalli et al. (2011). If the Hα emission is considered to dwarf galaxy VCC 1249 (Arrigoni Battaia et al. 2012)which originate from the same cylindrical volume of the X-ray gas, as- is, however, stripped by the ISM of an elliptical galaxy (M49), suming an electron temperature of 104 K and case B of nebular whereas IC 3418 is stripped by the ICM in the Virgo cluster.

A99, page 13 of 15 A&A 556, A99 (2013)

6 Hydrodynamic AMR simulations of Tonnesen & Bryan tail, we have found only (3σ) upper limits of <10 M in (2012) showed that star formation may occur in the ram pressure the three observed points covering the outer tail regions. The stripped tails when radiative cooling can overcome the heating estimated upper limit on the molecular gas content of the 7 from the surrounding ICM and thus is able to cool down the whole tail is ∼1 × 10 M. This is an amount that is similar stripped ISM. Our simplified calculations in Sect. 7 have sug- to the estimates based on the observed SFR averaged over gested that although molecular gas cannot be directly stripped by the tail, assuming a normal gas depletion timescale of 2 Gyr. ram pressure from IC 3418, some higher density H i than is usual Thus with our current sensitivity we cannot claim that the −2 in larger spirals in Virgo (ΣISM < 25 M pc within the inner relationship between the dense gas content and the SFR in −2 2 kpc disk radius, and <50 M pc in the outer disk radii) can be the tail is unusual. stripped thanks to a shallow potential of the galaxy and its prox- 3. Very sensitive upper limits on the 0.2–5 keV luminosity imity to the cluster center. It is then possible that these higher- of the gas-stripped tail has been obtained with Chandra. density clumps were easier to cool and compress by ram pressure The corresponding X-ray mass limit of the tail is <3.7 × 7 together with the thermal pressure of the ICM, supporting their 10 fX1/2 M (where fX is the filling factor of the X-ray emit- condensation. Moreover, because of the supersonic motion of ting gas) which is a factor of ∼280 less than in the tail of the IC 3418 through the ICM, the thermal pressure of the post-shock Norma cluster ram pressure stripped galaxy ESO 137-001. environment surrounding the stripped ISM within its Mach cone The absence of a bright X-ray tail may be mainly due to a is probably higher than the pre-shock thermal pressure, which low gas content in the tail, associated with its advanced evo- could also enhance condensation of the stripped material. In sim- lutionary state in which gas from the main body of the galaxy ulations, the stripped ISM fragments strongly behind the bow is no longer supplied to the tail. A rather low thermal pres- shock as a result of turbulence and Rayleigh-Taylor instability sure of the surrounding ICM in the Virgo cluster could also (e.g., Roediger et al. 2006). As shown, e.g., by Zhang et al. account for our non-detection. Existing AMR hydrodynamic (2012), the inner disks in nearby dwarf irregulars have propor- numerical simulations would suggest that at least some soft tionally more cool gas than the outer disks. Leroy et al. (2008) X-ray emission should be observable in the tail of IC 3418; indeed suggested that H2 dominates ISM of nearby dwarfs in however, these models do not study the evolutionary state the inner parts of their disks (<0.25r25) and not at the outer radii. after the galaxy is fully stripped. Then ram pressure stripped cool gas and denser clumps originate 4. Our semi-analytic calculations have indicated that preferentially from the inner disk regions. The global head-tail ISM parcels with rather high column densities up to −2 morphology of IC 3418 with a narrow (and linear) tail extending ∼25−50 M pc (depending on the galactic radius) can be behind a larger main body of the galaxy, would then be consis- accelerated from the galaxy. However, molecular clouds −2 tent with the proposed scenario. with densities >100 M pc cannot be directly stripped from the inner part of IC 3418 by the ram pressure in Virgo, 9. Conclusion even though this dwarf galaxy has a shallow galactic poten- tial, and it may get close to the cluster center. This explains We searched for cold molecular gas, as well as hot X-ray emit- why some molecular gas could survive in the galaxy center. ting gas, in IC 3418, a dwarf galaxy that is currently being trans- Moreover, the apparent loss of most of the molecular gas formed by ram pressure stripping from an irregular to an early- in the galaxy must be due to rather short lifetimes of dense type galaxy. It contains a prominent tail of young stars, a feature clouds that are similar to stripping timescales, and also due that is rare among other ram pressure stripped Virgo galaxies. to processes other than direct stripping, meaning some form Using the IRAM 30 m antenna, we searched for CO emission in of cloud destruction after the surrounding low density gas five deep integrations in the main body of the galaxy and in its is stripped. Such processes may contain ablation, which gas-stripped tail. New deep Chandra observations covered the includes Kelvin-Helmholtz instabilities, and self-destruction whole system. The main results are: of the clouds by supersonic motion of the galaxy through 1. A possible 12CO(2–1) 5σ marginal detection suggested that the Virgo ICM. 6 about 1.2 × 10 M of molecular gas (assuming a standard 5. The surviving central molecular gas reservoir makes it pos- Galactic CO-to-H2 conversion factor) is present in the cen- sible for a nuclear star cluster to eventually form, even after tral few arcseconds of IC 3418. Although CO(1–0) emission most of the galaxy’s gas has been stripped. was not formally detected in the center of the galaxy, there 6. We have suggested that ram pressure may act as an external is a 2σ feature with the same velocity range as the CO(2–1) agent that, together with the thermal pressure of the ICM and line, which is consistent with a weak emission diluted by assisted by radiative cooling, compresses the stripped dense the CO(1–0) large beam. An upper limit on molecular gas clumps and thus may trigger star formation in the tail of the content in the whole main body of the galaxy is estimated galaxy. 6 to <5 × 10 M, assuming the standard Galactic CO-to-H2 conversion factor. A similar limit is obtained from the ex- isting upper limits on FIR emission from dust. This is a factor of ∼20 less than the original molecular gas content Acknowledgements. We acknowledge support by the projects RVO:67985815, of the galaxy, as estimated from the UV-based SFR before M100031203 of the Academy of Sciences of the Czech Republic, GA CRˇ P209/11/P699, and GA CRˇ P209/12/1795. Pavel Jáchym was a Fulbright fel- quenching, and the typical gas depletion timescale in spirals low at Yale University in 2009–2010 when this project was prepared. Ming Sun of 2 Gyr. While some of this difference may arise from a non- is supported by NASA grants GO2-13102X, GO0-11145C, and GO1-12103A. standard CO-to-H2 conversion that could be ∼5 times larger We thank H. Crowl and S. Tonnesen for useful comments. The scientific results than Galactic, we think much of it is due to a true deficiency reported in this article are based in part on observations made by the Chandra / of molecular gas, given the lack of star formation in the main X-ray Observatory. This research has made use of the NASA IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California body. Institute of Technology, under contract with the National Aeronautics and Space 2. Although the presence of H ii regions in the tail suggests that Administration, and of the HyperLeda database (http://leda.univ-lyon1. there must be some dense (presumably molecular) gas in the fr).

A99, page 14 of 15 P. Jáchym et al.: Search for cold and hot gas in IC 3418

References Kennicutt, R. C., & Evans, N. J. 2012, ARA&A, 50, 531 Kewley, L. J., & Ellison, S. L. 2008, ApJ, 681, 1183 Abramson, A., Kenney, J. D. P., Crowl, H. H., et al. 2011, AJ, 141, 164 Kirby, E. M., Koribalski, B., Jerjen, H., & López-Sánchez, Á. 2012, MNRAS, Anders, E., & Grevesse, N. 1989, Geochim. Cosmochim. Acta, 53, 197 420, 2924 Arrigoni Battaia, F., Gavazzi, G., Fumagalli, M., et al. 2012, A&A, 543, A112 Komugi, S., Yasui, C., Kobayashi, N., et al. 2011, PASJ, 63, L1 Asplund, M., Grevesse, N., Sauval, A. J., & Scott, P. 2009, ARA&A, 47, 481 Koopmann, R. A., & Kenney, J. D. P. 2004, ApJ, 613, 851 Bell, E. F., & de Jong, R. S. 2000, MNRAS, 312, 497 Kormendy, J., & Bender, R. 2012, ApJS, 198, 2 Bell, E. F., & de Jong, R. S. 2001, ApJ, 550, 212 Leitherer, C., Schaerer, D., Goldader, J. D., et al. 1999, ApJS, 123, 3 Bigiel, F., Leroy, A., Walter, F., et al. 2008, AJ, 136, 2846 Leroy, A., Bolatto, A. D., Simon, J. D., & Blitz, L. 2005, ApJ, 625, 763 Binggeli, B., Tammann, G. A., & Sandage, A. 1987, AJ, 94, 251 Leroy, A., Bolatto, A., Stanimirovic, S., et al. 2007a, ApJ, 658, 1027 Binney, J., & Tremaine, S. 2008, Galactic Dynamics, 2nd edn. (Princeton Leroy, A., Cannon, J., Walter, F., Bolatto, A., & Weiss, A. 2007b, ApJ, 663, 990 University Press) Leroy, A. K., Walter, F., Brinks, E., et al. 2008, AJ, 136, 2782 Boissier, S., Boselli, A., Duc, P.-A., et al. 2012, A&A, 545, A142 Leroy, A. K., Walter, F., Sandstrom, K., et al. 2013, AJ, 146, 19 Bolatto, A. D., Leroy, A. K., Rosolowsky, E., Walter, F., & Blitz, L. 2008, ApJ, Lisenfeld, U., & Ferrara, A. 1998, ApJ, 496, 145 686, 948 Maloney, P., & Black, J. H. 1988, ApJ, 325, 389 Boselli, A., & Gavazzi, G. 2006, PASP, 118, 517 Matsushita, K., Belsole, E., Finoguenov, A., & Böhringer, H. 2002, A&A, 386, Boselli, A., Lequeux, J., & Gavazzi, G. 2002, A&A, 384, 33 77 Boselli, A., Boissier, S., Cortese, L., et al. 2005, ApJ, 623, L13 Mei, S., Blakeslee, J. P., Côté, P., et al. 2007, ApJ, 655, 144 Boselli, A., Boissier, S., Cortese, L., & Gavazzi, G. 2008, ApJ, 674, 742 Mizuno, N., Rubio, M., Mizuno, A., et al. 2001, PASJ, 53, L45 Bothwell, M. S., Kennicutt, R. C., & Lee, J. C. 2009, MNRAS, 400, 154 Mo, H., van den Bosch, F. C., & White, S. 2010, Galaxy Formation and Evolution Chung, A., van Gorkom, J. H., Kenney, J. D. P., & Vollmer, B. 2007, ApJ, 659, (Cambridge University Press) L115 Navarro, J. F., Frenk, C. S., & White, S. D. M. 1996, ApJ, 462, 563 Chung, A., van Gorkom, J. H., Kenney, J. D. P., Crowl, H., & Vollmer, B. 2009, Obreschkow, D., & Rawlings, S. 2009, MNRAS, 394, 1857 AJ, 138, 1741 Paturel, G., Petit, C., Prugniel, P., et al. 2003, A&A, 412, 45 Cortese, L., Gavazzi, G., Iglesias-Paramo, J., Boselli, A., & Carrasco, L. 2003, Pettini, M., & Pagel, B. E. J. 2004, MNRAS, 348, L59 A&A, 401, 471 Pineda, J. L., Goldsmith, P. F., Chapman, N., et al. 2010, ApJ, 721, 686 Cortese, L., Marcillac, D., Richard, J., et al. 2007, MNRAS, 376, 157 Poggianti, B. M., Smail, I., Dressler, A., et al. 1999, ApJ, 518, 576 Cortese, L., Bendo, G. J., Boselli, A., et al. 2010, A&A, 518, L63 Quilis, V., Moore, B., & Bower, R. 2000, Science, 288, 1617 Crowl, H. H., & Kenney, J. D. P. 2006, ApJ, 649, L75 Randall, S., Nulsen, P., Forman, W. R., et al. 2008, ApJ, 688, 208 Crowl, H. H., Kenney, J. D. P., van Gorkom, J. H., & Vollmer, B. 2005, AJ, 130, Roediger, E., & Hensler, G. 2005, A&A, 433, 875 65 Roediger, E., Brüggen, M., & Hoeft, M. 2006, MNRAS, 371, 609 Evans, A. S., Mazzarella, J. M., Surace, J. A., et al. 2005, ApJS, 159, 197 Schindler, S. 1999, A&A, 349, 435 Feldmann, R., Gnedin, N. Y., & Kravtsov, A. V. 2012, ApJ, 758, 127 Schruba, A., Leroy, A. K., Walter, F., et al. 2012, AJ, 143, 138 Fumagalli, M., & Gavazzi, G. 2008, A&A, 490, 571 Shibata, R., Matsushita, K., Yamasaki, N. Y., et al. 2001, ApJ, 549, 228 Fumagalli, M., Krumholz, M. R., Prochaska, J. X., Gavazzi, G., & Boselli, A. Smith, R. J., Lucey, J. R., Hammer, D., et al. 2010, MNRAS, 408, 1417 2009, ApJ, 697, 1811 Sun, M., & Vikhlinin, A. 2005, ApJ, 621, 718 Fumagalli, M., Gavazzi, G., Scaramella, R., & Franzetti, P. 2011, A&A, 528, Sun, M., Donahue, M., & Voit, G. M. 2007, ApJ, 671, 190 A46 Sun, M., Donahue, M., Roediger, E., et al. 2010, ApJ, 708, 946 Gavazzi, G., Boselli, A., Donati, A., Franzetti, P., & Scodeggio, M. 2003, A&A, Taylor, C. L., Kobulnicky, H. A., & Skillman, E. D. 1998, AJ, 116, 2746 400, 451 Tonnesen, S., & Bryan, G. L. 2009, ApJ, 694, 789 Gavazzi, G., Boselli, A., van Driel, W., & O’Neil, K. 2005, A&A, 429, 439 Tonnesen, S., & Bryan, G. L. 2010, ApJ, 709, 1203 Gavazzi, G., Boselli, A., Cortese, L., et al. 2006, A&A, 446, 839 Tonnesen, S., & Bryan, G. L. 2012, MNRAS, 2609 Ghigna, S., Moore, B., Governato, F., et al. 1998, MNRAS, 300, 146 Tonnesen, S., Bryan, G. L., & Chen, R. 2011, ApJ, 731, 98 Gil de Paz, A., Boissier, S., Madore, B. F., et al. 2007, ApJS, 173, 185 Tremonti, C. A., Heckman, T. M., Kauffmann, G., et al. 2004, ApJ, 613, 898 Guillard, P., Boulanger, F., Pineau des Forêts, G., et al. 2012, ApJ, 749, 158 Tuffs, R. J., Popescu, C. C., Pierini, D., et al. 2002, ApJS, 139, 37 Gunn, J. E., & Gott, III, J. R. 1972, ApJ, 176, 1 van Gorkom, J. H. 2004, Clusters of Galaxies: Probes of Cosmological Structure Hester, J. A., Seibert, M., Neill, J. D., et al. 2010, ApJ, 716, L14 and Galaxy Evolution, 305 Hunter, D. A., Elmegreen, B. G., & Ludka, B. C. 2010, AJ, 139, 447 Vollmer, B. 2009, A&A, 502, 427 Israel, F. P. 1997, A&A, 328, 471 Vollmer, B., Cayatte, V., Balkowski, C., & Duschl, W. J. 2001, ApJ, 561, 708 Jáchym, P., Palouš, J., Köppen, J., & Combes, F. 2007, A&A, 472, 5 Vollmer, B., Braine, J., Pappalardo, C., & Hily-Blant, P. 2008, A&A, 491, 455 Jáchym, P., Köppen, J., Palouš, J., & Combes, F. 2009, A&A, 500, 693 Vollmer, B., Wong, O. I., Braine, J., Chung, A., & Kenney, J. D. P. 2012, A&A, Kalberla, P. M. W., Burton, W. B., Hartmann, D., et al. 2005, A&A, 440, 775 543, A33 Kapferer, W., Sluka, C., Schindler, S., Ferrari, C., & Ziegler, B. 2009, A&A, We˙zgowiec, M., Vollmer, B., Ehle, M., et al. 2011, A&A, 531, A44 499, 87 Woudt, P. A., Kraan-Korteweg, R. C., Lucey, J., Fairall, A. P., & Moore, S. A. W. Kenney, J. D. P., & Koopmann, R. A. 1999, AJ, 117, 181 2008, MNRAS, 383, 445 Kenney, J. D., & Young, J. S. 1986, ApJ, 301, L13 Yoshida, M., Yagi, M., Komiyama, Y., et al. 2008, ApJ, 688, 918 Kenney, J. D. P., & Young, J. S. 1989, ApJ, 344, 171 Young, J. S., Xie, S., Tacconi, L., et al. 1995, ApJS, 98, 219 Kenney, J. D. P., van Gorkom, J. H., & Vollmer, B. 2004, AJ, 127, 3361 Zhang, H.-X., Hunter, D. A., & Elmegreen, B. G. 2012, ApJ, 754, 29

A99, page 15 of 15